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PHY 475/375 Lecture 10 (April 25, 2012) Multiple-Component Universes

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PHY 475/375 Lecture 10 (April 25, 2012) Multiple-Component Universes PHY 475/375 Lecture 10 (April 25, 2012) Multiple-Component Universes So far, we have studied model universes that contain only one component — matter only, radiation only, etc. In this chapter, we will look at more sophisticated models containing two or more components. We begin, as always, with the Friedmann equation; recall that it can be written in the form 2 2 8πG κc H(t) = 2 ε(t) 2 2 (6.1) 3c − R0 a(t) where H a/a˙ and ε(t) is the energy density contributed by all the components of the universe. ≡ Also, recall from an earlier lecture that equation (4.31) gave us a relation between κ, R0,H0, and Ω0: 2 κ H0 2 = 2 Ω0 1 (6.2) R0 c − Let us rewrite the Friedmann equation (6.1) without explicitly including the curvature by using equation (6.2): 8πG H2 H(t)2 = ε(t) 0 Ω 1 (6.3) 3c2 − a(t)2 0 − 2 Dividing by H0 , this becomes 2 H(t) ε(t) 1 Ω0 2 = + − 2 (6.4) H0 εc,0 a(t) where εc,0 is the critical density in the current epoch, given by 3c2H2 ε 0 (6.5) c,0 ≡ 8πG Now, we know that our Universe contains matter, for which the energy density εm has the depen- 3 4 dence εm = εm,0/a , and radiation, for which the energy density εr has the dependence εr = εr,0/a . It may also have a cosmological constant, with energy density εΛ = εΛ,0 = constant. While it is certainly possible that the universe contains other components as well, we will consider only the contributions of matter (w = 0), radiation (w =1/3), and the cosmological constant Λ (w = 1). − Recall from an earlier lecture that the energy density ε is additive, as written in equation (5.4): ε = εw w X PHY 475/375 (Spring 2012 ) Lecture 10 Using the additive property of the energy densities of different components, together with the various dependencies on the scale factor written above, equation (6.4) may be written as 2 H εr + εm + εΛ 1 Ω0 2 = + −2 H0 εc,0 a 4 3 εr,0/a + εm,0/a + εΛ,0 1 Ω0 = + −2 εc,0 a εr,0/εc,0 εm,0/εc,0 εΛ,0 1 Ω0 = 4 + 3 + + −2 a a εc,0 a 2 H Ωr,0 Ωm,0 1 Ω0 2 = 4 + 3 + ΩΛ,0 + −2 (6.6) ⇒ H0 a a a where we have put Ωr,0 = εr,0/εc,0, Ωm,0 = εm,0/εc,0, ΩΛ,0 = εΛ,0/εc,0, and Ω0 = Ωr,0 + Ωm,0 + ΩΛ,0. Note also that the Benchmark model introduced in an earlier lecture has Ω0 = 1 (i.e., spatially flat), but while this is consistent with the observational data, it is not demanded by the data. Therefore, we will retain the curvature term (1 Ω )/a2 in equation (6.6). − 0 To proceed, we write H =a/a ˙ in equation (6.6), multiply both sides by a2, and take the square root; this gives Ω Ω 1/2 H−1 a˙ = r,0 + m,0 + a2 Ω + (1 Ω ) (6.7) 0 a2 a Λ,0 − 0 Integrating, we get the cosmic time t as a function of a: a da = H t (6.8) [Ω /a2 + Ω /a + a2 Ω + (1 Ω )]1/2 0 Z0 r,0 m,0 Λ,0 − 0 In general, this integral does not have a simple analytic solution. There are instances, however, when simple analytic approximations may be found. For example, we discussed in an earlier lecture how in a multi-component universe containing radiation, matter, and Λ, the radiation dominates in the early stages. With radiation dominating, and Ω0 = 1 for a flat universe, equation (6.8) simplifies to a da 1 a 1 a2 a H0 t = a da = ≈ Ω /a2 Ω Ω 2 Z0 r,0 r,0 Z0 r,0 0 p p p so that we get eventually a2 H0 t (6.9) ≈ 2 Ωr,0 p Page 2 of 12 PHY 475/375 (Spring 2012 ) Lecture 10 From equation (6.9), we get 1/2 a(t) 2 Ω H t (6.10) ≈ r,0 0 p In the limit Ωr,0 = 1 (i.e., a spatially flat universe containing only radiation), this becomes 1/2 a(t) 2 H t ≈ 0 Using t0 = 1/2H0 from equation (5.62) that we wrote in the last lecture for a flat universe containing only radiation, this becomes t 1/2 a(t) ≈ t 0 which matches equation (5.64) that we obtained for the scale factor in a spatially flat universe containing only radiation. Likewise, there will be epochs where matter dominates, or Λ dominates, as we discussed in an earlier lecture. During some epochs, however, two of the components are of comparable density, and provide terms of roughly equal size in the Friedmann equation. During these epochs, a single-component model is a poor description of the universe, and we need to use a two-component model. We will now look at some of these instances in more detail. First, we will examine a universe which is of historical interest to cosmology: a universe containing both matter and curvature (either negative or positive). In the years following Einstein’s revoking of Λ, cosmologists studied in detail the possibility of a spatially curved universe dominated by non-relativistic matter. Matter + Curvature In a curved universe containing only matter, we can put Ωr,0 =0, ΩΛ,0 = 0. Also, since Ω0 = Ωr,0 + Ωm,0 + ΩΛ,0, having made the above choices (Ωr,0 =0, ΩΛ,0 = 0), we now have Ωm,0 = Ω0. With the above choices, the Friedmann equation (6.6) in a curved, matter-dominated universe can be written in the form 2 H(t) Ω0 1 Ω0 2 = 3 + −2 (6.12) H0 a a Note that although Ωm,0 = Ω0 is the only component, we retain the term (1 Ω0) for the curvature which means that Ω = 1; in other words, ε = ε . − 0 6 m,0 6 c,0 Page 3 of 12 PHY 475/375 (Spring 2012 ) Lecture 10 Let us now look at the characteristics of such a curved, matter-dominated universe. Suppose it is currently expanding (H0 > 0). 3 For it to stop expanding, we must have H0 = 0. Since Ω0/a is always positive, H(t)=0 • requires the second term on the right hand side of equation (6.12) to be negative. This means that a matter-dominated universe will cease to expand only if Ω0 > 1, and hence from equation (6.2), κ = +1 (positively curved). We can find the scale factor amax at maximum expansion by setting H(t) = 0 in equation • (6.12): Ω0 1 Ω0 0= 3 + −2 (6.13) amax amax so that a3 a2 max = max Ω Ω 1 0 0 − from which we get Ω a = 0 (6.14) max Ω 1 0 − Recall that Ω0 is the density parameter measured at a scale factor a(t0)=1. Also, note that in equation (6.12), the Hubble parameter enters only as H2, so there is a • time symmetry, and the contraction phase is just the time reversal of the expansion phase (although not at small scales, that is, small-scale processes will not be reversed during the contraction phase; e.g., stars will not absorb the photons they previously emitted). The eventual collapse of a Ω > 1 universe is sometimes called the “Big Crunch” — the • 0 universe returns to the hot, dense state in which it had originally started. In its contracting stage, an observer will see galaxies with a blueshift proportional to their distance, and wavelengths in the cosmic microwave background will compress progressively to shorter ones to eventually end in a cosmic γ-ray background. What happens, however, in a matter-dominated universe with Ω < 1, corresponding to κ = 1 0 − (negatively curved)? Both terms on the right hand side of equation (6.12) are then positive, so if such a universe • is expanding at t = t0, it will continue to expand forever (“Big Chill”). At early times, when the scale factor is small (a Ω0/ 1 Ω0 ), the matter term in the • Friedmann equation (6.12) will dominate, so we can≪ write{ fro−m equation} (6.8) that a a da a da a da a3/2 H t = = = 0 [Ω /a + (1 Ω )]1/2 [1/a + (1 Ω )/Ω ]1/2 ≈ [1/a]1/2 3/2 Z0 0 − 0 Z0 − 0 0 Z0 0 where we have used the fact that if a Ω0/ 1 Ω0 , then 1/a Ω0/ 1 Ω0 . From this expression, we see that the scale factor≪ will grow{ − at the} rate a ≫t2/3 at{ early− times.} ∝ Ultimately, the density of matter will be diluted far below the critical density, and the • universe will expand like the negatively curved empty universe we discussed in the last lecture, with a t. ∝ Page 4 of 12 PHY 475/375 (Spring 2012 ) Lecture 10 A plot of the scale factor a(t) vs. time, in units of H(t t0) is shown in Figure 6.1 of your text, and reproduced below. The panel on the right is an enlarged− view of the region enclosed by the small rectangle near “0” in the panel on the left. The solid line shows the fate of a matter-dominated universe with Ω0 > 1 (corresponding to • κ = +1, positively curved).
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